PROBING THE CORRELATION BETWEEN IR AND X-RAY EMISSION IN LUMINOUS INFRARED GALAXIES

Master thesis PROBING THE CORRELATION BETWEEN IR AND X-RAY EMISSION IN LUMINOUS INFRARED GALAXIES July 3, 2015 Author: Núria Torres Albà Advisor: Kazushi Iwasawa Tutor: Valentí Bosch Ramon Universitat de Barcelona Dept. d Astronomia i Meteorologia

Acknowledgement I would like to thank my advisor, Kazushi Iwasawa, for all the help and guidance provided throughout all these months. His extensive knowledge and experience on the field of study has been crucial for this work s development, providing me with a proper understanding of the subject. I would also like to thank my tutor, Valentí Bosch-Ramon, for his counselling and direction in the final theoretical development, as well as his advice and revisions. In addition, I also thank Josep Maria Paredes for the final comments and suggestions, helping in the improvement of this work s clarity. 2

Abstract In this work, X-ray data for a sample of 59 luminous infrared galaxies (LIRGs), obtained with the Chandra X-ray Observatory is presented. These are observations of the lower luminosity portion of the Great Observatory All-sky LIRG Survey (GOALS), which includes the most luminous infrared selected galaxies, log(l ir /L ) 11.05, in the local universe, z 0.088.With combined X-ray and mid infrared diagnostics for AGN, AGN are found in 32 % of the objects. Objects without traces of AGN appear to be underluminous in X-rays, compared to the previously derived far infrared and X-ray correlations for starburst galaxies at lower star formation rates. Our results suggest the previous claim that star formation rate can be directly inferred from X-ray luminosity may not hold for dusty galaxies like LIRGs. 3

Contents 1 Introduction 5 1.1 Luminous infrared galaxies (LIRGs).................. 5 1.1.1 Origin of the infrared emission................. 5 1.1.2 Merger: Origin of AGN and starburst............. 6 1.2 X-ray study................................ 7 1.2.1 The X-ray spectrum of an AGN................ 7 1.2.2 The X-ray spectrum of a starburst region........... 8 2 X-ray observations 10 2.1 Chandra and XMM-Newton....................... 11 3 GOALS 12 4 Chandra observations of LIRGs 12 4.1 The sample................................ 12 4.2 Observations and data reduction.................... 13 4.2.1 Hard band fitting......................... 16 4.2.2 Soft band fitting......................... 16 4.2.3 Special cases........................... 17 4.3 Results................................... 18 4.3.1 AGN selection.......................... 18 4.3.2 X-ray luminosities and correlation with IR........... 20 5 Discussion 24 5.1 X-ray and infrared luminosity correlation............... 24 5.2 Molecular disk and AGN interaction.................. 26 4

1 Introduction In astronomy, many methods are used to estimate the star formation rate in galaxies, as described in a review by Kennicutt (1998), the basis of which are different manners of quantifying the amount of hot, young, and short-lived stars. One of these methods is a measurement of the energy deposed in the dust of the interstellar medium by these stars, which is obtained through far-infrared luminosity measurements. Therefore, far-infrared luminosity can be directly related to the star formation rate, and is probably the best method to determine it. In galaxies with a considerable amount of star formation, such as starburst galaxies, emission in other wavelengths can be somewhat related to young and massive stars, such as X-ray luminosity (e.g. X-ray binaries emission, supernova remnants (SNRs)). Therefore, it has been suggested that if a good correlation between X-ray luminosity and IR luminosity exists in galaxies, the star formation rate (SFR) can be directly inferred from the X-ray luminosity (e.g. Ranalli et al., 2003; Grimm et al., 2003; Mineo et al., 2014). However, the applicability of this method has been questioned due to experimental results at higher SFRs (e.g. Iwasawa et al., 2009). The aim of this work is the study of this correlation through the calculation of the X-ray luminosity derived from Chandra spectral data in a subsample of GOALS galaxies, in order to further constrain its usability. 1.1 Luminous infrared galaxies (LIRGs) Luminous infrared galaxies (LIRGs) are galaxies with an infrared luminosity between 10 11 L and 10 12 L, emitting more in this band than in all other bands combined. Galaxies with an even higher infrared luminosity, above 10 12 L are called ULIRGS, Ultra Luminous Infrared Galaxies. 1.1.1 Origin of the infrared emission Galaxies emit in the infrared mainly through three sources: stars, interstellar gas and dust. The emission from starts peaks in the near infrared (1-3 microns), and the emission from atoms and molecules in the interstellar gas makes up only a few percent of the total infrared luminosity in galaxies. Dust is heated by startlight and emits thermally, clearly being the dominating infrared source beyond 3 microns. The reason some galaxies emit more in the infrared than others is due to a high increase in the dust emission, which is due to an increase in the amount of energetic photons coming from the sources that heat this dust. As dust absorbs preferentially UV emission, since the cross-section is highest for UV light, most of the absorbed photons will come from the massive O and B stars. Another important contribution to the dust heating is the continuum emission from an active nuclei s accretion disk, as it peaks in the far UV. Therefore, the high infrared luminosity in LIRGs may come from both an ex- 5

tremely high star formation rate, which is known as starburst, giving raise to many young luminous stars, and the contribution of an active galactic nuclei (AGN) component. 1.1.2 Merger: Origin of AGN and starburst When cool gas is available in a galaxy it settles in a disk-like morphology to conserve angular momentum. As this configuration is stable, there is need of a merger in order to funnel it towards the central regions of a galaxy, where it can feed both star formation and the central black hole, thus both processes being somewhat correlated. Kauffman et al. (2003a) found that low-luminosity AGN rather live in galaxies with older stellar populations, while the high-luminosity ones tend to be hosted by galaxies with young populations. In fact, strong Hδ absorption lines in highluminosity sources indicate that these galaxies have experienced a starburst some 0.1-1 Gyr ago, and tracers of starburst tend to correlate with AGN activity indicators (e.g. Schweitzer et al., 2006). Silverman et al. (2009) studied the hosts of 147 Chandra detected AGN, which appear to preferentially be inside environments with substantial star formation. The same conclusion was reached by Rumbaugh et al. (2012), finding that on average AGN hosts either are still active starburst galaxies or have experienced one within the previous 1 Gyr. However, AGN and star formation processes do not necessarily end at the same time: the mass involved in an active star formation phase is some hundreds to thousand of times larger than the mass accreted onto the black hole. Therefore, once the starburst dies out, AGN activity can continue. As the galaxy is more and more depleted of gas, even the black hole s activity comes to a halt, entering a quiescent stage. It is believed that all galaxies, now known to host a supermassive central black holes, have once gone through both a period of strong activity and a high star formation rate. Therefore, the study of these starburst galaxies can significantly contribute to the understanding of galaxy evolution. Another key point derived from the study of these type of galaxies is galaxy building, since more than 95% of ULIRG sources are mergers (Sanders and Mirabel, 1996). The current favoured model for the building up is hierarchical growth, which is based on the idea that galaxies form mainly through the merging of nearby smaller galaxies to give rise to larger structure, which will then undergo further merging. Sanders (1999) thus puts the ULIRGs as a stage or episode between gas-rich merging disk galaxies and the resulting ellipticals. Both cosmological simulations and observation of large scale structure, which becomes more pronounced towards lower redshifts, support this theory. In more than 70% of the cases, AGN host galaxies have nearby companions, some showing tidal tails, which are strong evidence for interaction (Beckmann and Shrader, 2012). These tidal forces can, prior to the merger per se, lead to a distur- 6

bance, which allows matter to flow into the centre towards the black hole, as well as triggering star formation in the galaxies due to compression of gas. After a major merger, most of the cold gas is driven to the core of the galaxies remnant and this large in-flow of gas can feed or trigger these startbust and AGN processes. If an efficient mass intake takes place, the SFR and black hole growth accelerate rapidly, within a few 10 8 years. Based on simulations of the AGN activation through major merging events, Hopkins et al. (2005) derive a black hole growth time scale of about 100 Myr. The AGN will be active throughout this time, but surrounded by obscuring material, and thus will appear dim. Most of the accretion would occur during this obscured phase, which would end with the clearing of the surrounding material, probably due to the AGN s own activity (heating of the surrounding gas, expulsion due to strong winds or jets...), and the natural depletion of gas after being consumed. Afterwards, the object will appear as a bright quasar for a duration of 10 Myr, before exhausting completely its fuel and returning to a quiescent state (Sanders, 1999). 1.2 X-ray study Both the AGN and starburst component in a LIRG are observed in a large range of wavelengths, from radio to high energies. Electromagnetic radiation between 120 ev and 120 kev, associated to thermal temperatures above 10 7 K, is referred to as X-rays. The study of radiation in this domain has seen an enormous evolution over recent decades, and is currently one of the key energy ranges to study AGN. X-ray photons have low extinction cross section in the ISM, therefore X-ray observations are much less vulnerable to dust extinction than optical or UV observations. Chandra observations routinely penetrate A V > 100 mag of optical extinction (Wang, 2007). Therefore, in order to study such highly obscured galaxies, X-ray emission is particularly adequate. Also, AGN are the dominating sources of cosmic X-rays: in any medium-deep image of the X-ray sky at high galactic latitudes, the majority of detections (about 80%, Gandhi, 2005) are AGN. 1.2.1 The X-ray spectrum of an AGN X-rays are the most direct probe of the accretion processes responsible for such high energy emission, as they can penetrate through large columns of obscuring gas and dust. Accreting matter falls towards the black hole, settling in a planar orbit around it, where gravitational energy is released through dissipation. Temperature increases in the disk when approaching the central engine, reaching the X-ray generating temperatures at a few gravitational radii from it, producing a thermal spectrum. The low-energy photons produced are scattered through inverse Compton to higher energies by relativistic electrons in the corona (Haardt and Maraschi, 1993). The 7

inverse Compton spectrum has a power law shape with a photon index of Γ 2, and as the disk temperature and the electron energies are limited, it has a high energy cutoff. Superposed to this powerlaw, reflection of photons from the corona leads to the appearance of a hump in the hard X-ray spectrum, generally peaking at around 20 30 kev, where the reflection efficiency has its maximum. Another signature of the reflection process is the 26 F ek α1 fluorescence line at E Kα = 6.404 kev. A representation of the resulting spectrum can be seen in Fig. 1. The origin of reflection is theorized to be caused by a cold, optically-thick material. As X-ray continuum illuminates the material, photons can either be Compton scattered by free or bound electrons or they can be absorbed, followed by fluorescent line emission (George and Fabian, 1991). The origin of the soft excess emission is still an open issue (Dewangan et al., 2007; Crummy et al., 2006; Done, 2007) Figure 1: Schematic representation of a type 1 AGN spectrum in the X-ray. Image from Ricci (2012). AGN are classified into type 1 or type 2 based on the measurement of intrinsic absorption in the X-ray band. The lower energy spectrum of an AGN (below 10 kev) is modified by obscuring gas through photoelectric absorption if the line of sight passes through the obscuring torus. At a column density of N H 1.5 10 24 cm 2 the energy transmitted at 10 kev decreases one e-folding, being these sources called Compton-thick. At higher column densities, multiple scatterings, as modelled by the Klein-Nishina cross section, lead to a Compton down-scattering over the full X-ray regime. At column densities of about N H 1.5 10 24 cm 2, barely any flux emerges (Fig. 2), as all photons are scattered more than once. 1.2.2 The X-ray spectrum of a starburst region An enhanced star formation rate comes with an increase of the energetic phenomena associated with the final stages of stellar evolution: X-ray binaries, supernova remnants (SNRs), galactic winds and, it has been predicted though still not ob- 8

Figure 2: Modelled AGN rest-frame X-ray spectral energy distributions, showing intrinsic power law emission, reflection effects and high energy exponential cut-off. Obscuration by material is presented for different column densities, labelled as log[n H] = 20.25 to 25.25, in units of atoms cm 2. The shaded region marks the X-ray spectral regime that is accessible for study to current high spatial resolution X-ray instruments such as Chandra and XMM-Newton. Figure from Gandhi (2005). served, Compton scattering of ambient FIR photons by SN-accelerated relativistic electrons (for a more detailed description, refer to Persic and Rephaeli (2002)). A typical starburst spectrum (as the one shown in Fig. 3) has the following main contributions: Binary systems are the brightest galactic X-ray sources, and it is generally assumed that they dominate the emission in starburst galaxies in the range of 2 10 kev. The high-mass X-ray binaries (HMXB) are formed by a neutron star or a black hole and an optical companion of M opt 8M. X-ray emission comes from the accretion by the massive compact object of matter driven by the high mass star s wind. This generates a powerlaw spectra of photon index Γ 1.2, so harder than that of a typical AGN. A striking characteristic in the HMXB spectra is a prominent Fe-K emission feature with central energy in the range 6.4 6.7 kev, generated in the companion s stellar wind or in the compact object s magnetosphere. Massive stars, with M 8M, will end up exploding as supernovae. X-ray emission from this sources comes mainly from the first phase of evolution of the SNR, the free expansion, as its material shocks and sweeps the surrounding ISM and heats gas. This results in a thermal spectrum with kt 2 kev, coming primarily from bremsstrahlung continuum and collisionally excited lines, one of which can be the mentioned intense Fe-K line. A contribution of nonthermal emission may also be present, coming from the central pulsar, or from 9

synchrotron emission by electrons accelerated in the shock. The high SN rate in a starburst implies an increase in relativistic electron densities, as SN shocks are important sites of cosmic ray acceleration. These relativistic electrons can Compton up-scatter low energy FIR photons, abundant in a starburst region as previously explained, to hard X-ray energies, resulting in a powerlaw spectrum with the photon index of the electron population. An important contribution in the softer X-ray emission is that coming from more external ISM and the galactic halo gas, shock-heated by the multiple SN explosions and stellar winds, with kt 1 kev. Thermal soft X-ray emission is therefore expected from this gas. As all these contributions are directly related to the star formation, in starburst regions X-ray luminosity can be related to the star formation rate. Figure 3: XMM-Newton spectrum of galaxy M82 (K. Iwasawa, private communication), typical starburst, with multiple emission lines and a hard tail (left). X-ray image of M82 (right, Chandra photo album: January 13, 2011.). Red: 0.3 1.1 kev, Green: 0.7 2.2 kev, Blue: 2.2 6 kev. Resolved point sources are X-ray binaries. 2 X-ray observations Unlike modern optical astronomy, which began about 400 years ago with the use of the first telescopes, X-ray astronomy is only about four decades old. The main reason for this late development is the high opaqueness of the atmosphere to X-rays, due to the high photoelectric absorption cross-section of this energetic photons in air. The second reason is that the very small wavelengths at such high energies require new ways to focus light in order to form images. The first of these two problems implies the necessity of performing X-ray observations in space, with the use of artificial satellites. 10

The second difficulty arises from the fact that X-ray photons typically have energies 1000 times higher than that of optical photons. Thus, normal incidence mirrors generally used to focus optical light completely absorb X-rays, making it impossible to use the well known telescope techniques in optical regime. To focus X-rays, two conditions are necessary. The first one is for the incidence angle to be very small, in order to avoid absorption of the photons. This reflection angle varies as ρ/e, where ρ is the density of the reflecting surface and E the energy of the incoming X-ray photon (Aschenbach, 1985). Thus, the higher this energy, the smaller the angle has to be in order to obtain reflection. For photons of a few kev, this angle turns out to be smaller than 1 o. The second requirement is that the reflecting surface has to be highly polished, to an average roughness of only a few angstroms, given the small reflection angles and photon wavelengths. Alignment, polishing and coating of the mirror surfaces are the most critical parts of the assembly. Therefore, X-ray mirrors are highly polished surfaces placed almost parallel to the incoming beam. The reflection of a paraboloid section followed by a confocal and coaxial hyperboloid surface provides good focus on-axis and also at a small off-axis angles (about a few arcmin at most), as seen in Fig.4. As the mirrors are placed in the direction of the incoming light, the effective surface of reflection is very small, and thus the sensitivity is low. To solve this, the reflecting area is increased by nesting concentric surfaces of different radii. Figure 4: Main elements of a Wolter type 1 nested X-ray focusing mirror configuration. Image from http://chandra.harvard.edu/. 2.1 Chandra and XMM-Newton Two major observatories were launched in 1999: NASA s Chandra satellite and ESA s XMM-Newton mission. Chandra has four pairs of iridum-coated mirrors, the largest one with a diameter of 120 cm. Depending on the wavelength, the effective area ranges from 800 cm 2 11

to 40 cm 2, being the lower value for the shortest wavelengths. Its focal length is of 10 m. Chandra s best feature is that it allows an incredibly good spatial resolution, of 0.5 arcsec, the best of any X-ray instrument ever built. Its observation range is from 0.1 to 10 kev. XMM Newton is comprised of 3 X-ray telescopes, each consisting of 58 nested mirrors coated with gold, the biggest of which has a diameter of 70 cm. Its focal length is of 7.5 m, and the effective area ranges from 1900 cm 2 at low energies to 350 cm 2 at high energies. It complements Chandra by having a much larger collecting area, but a worse spatial resolution by a factor of a few. It is able to search for spectral features in distant cosmic sources fainter than those detectable by any previous mission. 3 GOALS The Great Observatories All-sky LIRG Survey (GOALS, Armus et al., 2009), is combining imaging and spectroscopic data from NASA s Spitzer, Hubble, Chandra and GALEX space-borne observatories in a comprehensive study of over 200 of the most luminous infrared-selected galaxies in the local Universe. The sample consists of approximately 180 LIRGs as well as over 20 ULIRGs, 77 of which are systems that contain multiple galaxies. The objects are a complete subset of the IRAS Revised Bright Galaxy Sample (RBGS, Sanders et al., 2003), which comprises 629 extragalactic objects with 60-micron flux densities above 5.24 Jy and Galactic latitude above five degrees. The RBGS objects, all with redshifts, z < 0.088, are the brightest 60-micron sources in the extragalactic sky. The LIRGs and ULIRGs targeted in GOALS span the full range of nuclear spectral types (type-1 and type- 2 AGN, LINERs, and starbursts) and interaction stages (major mergers, minor mergers, and isolated galaxies). They provide an unbiased picture of the processes responsible for enhanced infrared emission in the local Universe, and are excellent analogues for comparisons with infrared and sub-millimetre selected galaxies at highredshift. 4 Chandra observations of LIRGs 4.1 The sample The sample studied in this work consists of 59 systems from the GOALS sample, 15 of which contain multiple galaxies. It represents the low luminosity part of the GOALS sample, with 11.05 log(l ir /L ) 11.73, incomplete as Chandra data for all systems is unavailable. The redshift range is z = 0.003 0.037. The median IR luminosity of the sample is log(l ir /L ) = 11.40, corresponding to log(l ir ) = 44.99. Table 1 gives basic data for all 59 systems in the sample. A complete study of the high luminosity part, 11.73 log(l ir /L ) 12.57, with 12

44 systems, was carried out by Iwasawa et al. (2011), with z = 0.010 0.088 and median log(l ir /L ) = 11.99, corresponding to log(l ir ) = 45.58. A comparison of the properties of the two samples is shown in Fig. 5. Figure 5: Luminosity distance (left) and infrared luminosity (right) distributions of the sample of study and the 44 system sample from Iwasawa et al., (2011). 4.2 Observations and data reduction Twenty-seven of the 59 objects were observed with Chandra in cycle-13 with a uniform 15 ks exposure on each target. All the observations were carried out in imaging mode with the ACIS-S detector operated in VFAINT mode. The Chandra data for the remaining 32 targets were taken from the Archive. The exposure times for these data go from 4.88 to 58.34 ks, all taken with ACIS-S in FAINT or VFAINT mode. When more than one observation was available, the one with longest exposure time was used. The observation log is shown in Table 2. Galactic HI column density is estimated from the Leiden/Argentine/Bonn (LAB) survey (Kalberla et al., 2005). The data reduction was performed using the Chandra data analysis package CIAO version 4.7, and HEASARC s FTOOLS. Except for some sources, which are clearly absorbed AGN, Chandra spectra of the sources appears similar: an emission-line dominated soft X-ray band and a hard X-ray tail. Examples of both spectra can be seen in Fig. 6. In order to estimate the X-ray fluxes for all sources, all components in multiple systems are analysed separately, as specified in Table 3. Fitting is done separately for the soft (0.5 2 kev) and hard (2 7 kev) X-ray components. Flux and luminosity in Table 3 are corrected only for galactic absorption, even if significant amount of intrinsic absorption may be present in the source. X-ray spectra of our objects is complex, and most of the spectra do not have sufficient quality to decompose and estimate intrinsic absorption for each component. For this reason, it is not corrected. 13

Table 1: The sample. No. (1) IRAS Name (2) Optical ID (3) RA (NED) (J2000) (4) DEC (NED) (J2000) (5) z (km s 1 ) (6) D L (Mpc) (7) log(l ir ) (L ) (8) 45 F13182+3424 UGC 08387 13 h20 m35.34 s +34 d08 m22.2 s 0.023299 110.0 11.73 47 F01173+1405 CGCG 436-030 01 h20 m02.72 s +14 d21 m42.9 s 0.031229 134.0 11.69 49 F01484+2220 NGC 0695 01 h51 m14.24 s 22 d34 m56.5 s 0.032472 139.0 11.69 50 F12592+0436 CGCG 043-099 13 h01 m50.80 s +04 d20 m00.0 s 0.037483 175.0 11.68 51 F11011+4107 MCG+07-23-019 11 h03 m53.20 s +40 d50 m57.0 s 0.034524 158.0 11.62 53 F02512+1446 UGC 02369 02 h54 m01.78 s +14 d58 m24.9 s 0.031202 136.0 11.67 54 F04315-0840 NGC 1614 04 h33 m59.85 s -08 d34 m44.0 s 0.015938 67.8 11.65 56 F13497+0220 NGC 5331 13 h52 m16.29 s +02 d06 m17.0 s 0.033043 155.0 11.66 57 F06076-2139 IRAS F06076-2139 06 h09 m45.81 s -21 d40 m23.7 s 0.037446 165.0 11.65 60 F11231+1456 IRAS F11231+1456 11 h25 m47.30 s +14 d40 m21.1 s 0.034167 157.0 11.64 63 18090+0130 IRAS 18090+0130 18 h11 m35.91 s +01 d31 m41.3 s 0.034167 134.0 11.65 64 F01417+1651 III Zw 035 01 h44 m30.45 s +17 d06 m05.0 s 0.027436 119.0 11.64 65 F10257-4339 NGC 3256 10 h27 m51.27 s -43 d54 m13.8 s 0.009354 38.9 11.64 67 F16399-0937 IRAS F16399-0937 16 h42 m40.21 s -09 d43 m14.4 s 0.027012 128.0 11.63 68 F16164-0746 IRAS F16164-0746 16 h19 m11.79 s -07 d54 m02.8 s 0.027152 128.0 11.62 71 F08354+2555 NGC 2623 08 h38 m24.08 s +25 d45 m16.6 s 0.018509 84.1 11.60 72 F23135+2517 IC 5298 23 h16 m00.70 s +25 d33 m24.1 s 0.027422 119.0 11.60 73 20351+2521 IRAS 20351+2521 20 h37 m17.72 s +25 d31 m37.7 s 0.033697 151.0 11.61 75 F16104+5235 NGC 6090 16 h11 m40.70 s +52 d27 m24.0 s 0.029304 137.0 11.58 79 F13362+4831 NGC 5256 13 h38 m17.52 s +48 d16 m36.7 s 0.027863 129.0 11.56 80 F03359+1523 IRAS F03359+1523 03 h38 m46.70 s +15 d32 m55.0 s 0.035401 152.0 11.55 81 F04210-4042 ESO 550-IG025 04 h21 m20.02 s -18 d48 m47.6 s 0.032196 135.8 11.51 82 F00085-1223 NGC 0034 00 h11 m06.55 s -12 d06 m26.3 s 0.019617 84.1 11.49 83 F00506+7248 MCG+12-02-001 00 h54 m03.61 s +73 d05 m11.8 s 0.015698 69.8 11.50 85 F17138-1017 IRAS F17138-1017 17 h16 m35.79 s -10 d20 m39.4 s 0.017335 84.0 11.49 95 F12043-3140 ESO 440-IG058 12 h06 m51.82 s -31 d56 m53.1 s 0.023203 112.0 11.43 100 F21453-3511 NGC 7130 21 h48 m19.50 s -34 d57 m04.7 s 0.016151 72.7 11.42 104 F23488+1949 NGC 7771 23 h51 m24.88 s +20 d06 m42.6 s 0.014267 61.2 11.40 105 F23157-0441 NGC 7592 23 h18 m22.20 s -04 d24 m57.6 s 0.024444 106.0 11.40 106 F16577+5900 NGC 6286 16 h58 m31.38 s +58 d56 m10.5 s 0.018349 85.7 11.37 107 F12590+2934 NGC 4922 13 h01 m24.89 s +29 d18 m40.0 s 0.023586 111.0 11.38 110 F10015-0614 NGC 3110 10 h04 m02.11 s -06 d28 m29.2 s 0.016858 79.5 11.37 114 F00402-2349 NGC 0232 00 h42 m45.82 s -23 d33 m40.9 s 0.022639 95.2 11.44 117 F09333+4841 MCG+08-18-013 09 h36 m37.19 s +48 d28 m27.7 s 0.025941 117.0 11.34 120 F15107+0724 CGCG 049-057 15 h13 m13.09 s +07 d13 m31.8 s 0.012999 65.4 11.35 121 F02401-0013 NGC 1068 02 h42 m40.71 s -00 d00 m47.8 s 0.003793 15.9 11.40 123 F02435+1253 UGC 02238 02 h46 m17.49 s +13 d05 m44.4 s 0.021883 92.4 11.33 127 F13197-1627 MCG-03-34-064 13 h22 m24.46 s -16 d43 m42.9 s 0.016541 82.2 11.28 134 F00344-3349 ESO 350-IG038 00 h36 m52.25 s -33 d33 m18.1 s 0.020598 89.0 11.28 136 F23394-0353 MCG -01-60-022 23 h42 m00.85 s -03 d36 m54.6 s 0.023236 100.0 11.27 142 F13229-2934 NGC 5135 13 h25 m44.06 s -29 d50 m01.2 s 0.013693 60.9 11.30 141 F13126+2453 IC 0860 13 h15 m03.53 s +24 d37 m07.9 s 0.011164 56.8 11.14 147 F22132-3705 IC 5179 22 h16 m09.10 s -36 d50 m37.4 s 0.011415 51.4 11.24 148 F03514+1546 CGCG 465-012 03 h54 m16.08 s +15 d55 m43.4 s 0.022222 94.3 11.20 163 F12243-0036 NGC 4418 12 h26 m54.62 s -00 d52 m39.2 s 0.007268 36.5 11.19 169 F21330-3846 ESO 343-IG013 21 h36 m10.83 s -38 d32 m37.9 s 0.019060 85.8 11.14 170 F06107+7822 NGC 2146 06 h18 m37.71 s +78 d21 m25.3 s 0.002979 17.5 11.12 174 F14280+3126 NGC 5653 14 h30 m10.42 s +31 d12 m55.8 s 0.011881 60.2 11.13 178 F12116+5448 NGC 4194 12 h14 m09.47 s +54 d31 m36.6 s 0.008342 43.0 11.10 179 F23157+0618 NGC 7591 23 h18 m16.28 s +06 d35 m08.9 s 0.016531 71.4 11.12 182 F00073+2538 NGC 0023 00 h09 m53.41 s +25 d55 m25.6 s 0.015231 65.2 11.12 188 F23133-4251 NGC 7552 23 h16 m10.77 s -42 d35 m05.4 s 0.005365 23.5 11.11 191 F04118-3207 ESO 420-G013 04 h13 m49.69 s -32 d00 m25.1 s 0.011908 51.0 11.07 194 08424-3130 ESO 432-IG006 08 h44 m28.07 s -31 d41 m40.6 s 0.016165 74.4 11.08 195 F05365+6921 NGC 1961 05 h42 m04.65 s +69 d22 m42.4 s 0.013122 59.0 11.06 14

No. (1) Table 1 continued. IRAS Name (2) Optical ID (3) RA (NED) (J2000) (4) DEC (NED) (J2000) (5) z (km s 1 ) (6) D L (Mpc) (7) log(l ir ) (L ) (8) 196 F23444+2911 Arp 86 23 h47 m01.70 s +29 d28 m16.3 s 0.016161 73.6 11.07 198 F03316-3618 NGC 1365 03 h33 m36.37 s -36 d08 m25.4 s 0.005457 17.9 11.00 199 F10196+2149 NGC 3221 10 h22 m19.98 s +21 d34 m10.5 s 0.013709 65.7 11.09 201 F02071-1023 NGC 0838 02 h09 m38.58 s -10 d08 m46.3 s 0.012844 53.8 11.05 Notes: Column(1): through number of the object, also used in other tables. Column (2): original IRAS source, where an F prefix indicates the Faint Source Catalog and no prefix indicates the Point Source Catalog. Column (3): optical cross-identification, where available from NED. Column (4): the best available right ascension (J2000) in NED as of October 2008. Column (5): the best available source declination (J2000) from NED as of October 2008. Column (6): the best available heliocentric redshift from NED as of October 2008. Column (7): the luminosity distance in megaparsecs derived by correcting the heliocentric velocity for the 3-attractor flow model of Mould et al. (2000) and adopting cosmological parameters H 0 = 73 km s 1 Mpc 2, Ω V = 0.73, and Ω M = 0.27, as provided by NED. Column (8): the total infrared luminosity in log 10 Solar units computed using the flux densities reported in the RGBS and the luminosity distances in Col. (7) using the formulae L ir /L = 4π(D L[m] ) 2. Table 2: Chandra observation log. No. Galaxy Obs ID Date Mode Exp. Time (ks) 0.5-7.0 kev a (cts) N H,Gal b (10 20 cm 2 ) 45 UGC 08387 7811 2007-02-19 VFAINT 14.07 277.9 ± 17.1 1.0 47 CGCG 436-030 15047 2012-11-24 VFAINT 13.82 168.6 ± 13.8 3.4 49 NGC 0695 15046 2013-01-01 VFAINT 14.78 312.9 ± 18.5 6.9 50 CGCG 043-099 15048 2012-11-23 VFAINT 14.78 71.9 ± 8.1 1.9 51 MCG+07-23-019 12977 2011-02-07 VFAINT 52.34 506.9 ± 26.8 1.0 53 UGC 02369 4058 2002-12-14 FAINT 9.68 120.6 ± 12.0 7.9 54 NGC 1614 15050 2012-11-21 VFAINT 15.76 800.0 ± 28.9 6.3 56 NGC 5331 15051 2013-05-12 VFAINT 14.78 121.9 ± 12.4 2.0 57 IRAS F06076-2139 15052 2012-12-12 VFAINT 14.78 52.4 ± 8.2 7.6 60 IC 2810 15053 2013-10-27 VFAINT 14.78 93.2 ± 11.7 2.5 63 IRAS 18090+0130 15054 2013-02-10 VFAINT 14.77 98.9 ± 11.3 20.2 64 III Zw 035 6855 2006-02-24 FAINT 14.98 81.4 ± 9.0 4.8 65 NGC 3256 835 2000-01-05 FAINT 27.80 8117.2 ± 102.3 9.1 67 IRAS F16399-0937 15055 2013-06-30 VFAINT 14.87 161.9 ± 14.4 13.0 68 IRAS F16164-0746 15057 2013-01-19 VFAINT 14.78 99.2 ± 11.3 11.3 71 NGC 2623 4059 2003-01-03 FAINT 19.79 171.0 ± 14.1 3.1 72 IC 5298 15059 2013-02-04 VFAINT 14.78 222.8 ± 16.0 5.7 73 IRAS 20351+2521 15058 2012-12-13 VFAINT 13.56 146.8 ± 14.0 13.1 75 NGC 6090 6859 2006-05-14 FAINT 14.79 347.5 ± 19.3 1.6 79 NGC 5256 2044 2001-11-02 FAINT 19.69 1451.2 ± 43.5 1.7 80 IRAS F03359+1523 6856 2005-12-17 FAINT 14.76 108.2 ± 11.4 13.8 81 ESO 550-IG025 15060 2012-11-24 VFAINT 14.78 72.2 ± 10.6 3.2 82 NGC 0034 15061 2013-06-05 VFAINT 14.78 329.0 ± 19.5 2.1 83 MCG+12-02-001 15062 2012-11-22 VFAINT 14.31 311.0 ± 19.3 22.0 85 IRAS F17138-1017 15063 2013-07-12 VFAINT 14.78 207-3 ± 15.4 17.0 95 ESO 440-IG058 15064 2013-03-20 VFAINT 14.78 187.0 ± 16.1 5.6 100 NGC 7130 2188 2001-10-23 FAINT 38.64 3327.1 ± 59.3 1.9 104 NGC 7771 10397 2009-05-22 VFAINT 16.71 904.6 ± 34.6 4.0 105 NGC 7592 6860 2006-10-15 FAINT 14.99 388.7 ± 21.9 3.8 106 NGC 6286 10566 2009-09-18 FAINT 14.00 544.8 ± 27.9 1.8 107 NGC 4922 15065 2013-11-02 VFAINT 14.86 202.9 ± 17.2 0.9 110 NGC 3110 15069 2013-02-02 VFAINT 14.87 396.3 ± 22.3 3.5 114 NGC 0232 15066 2013-01-04 VFAINT 14.78 193.5 ± 15.7 1.4 117 MCG+08-18-013 15067 2013-06-03 VFAINT 13.79 101.7 ± 11.1 1.7 15

Table 2 continued. No. Galaxy Obs ID Date Mode Exp. Time (ks) 0.5-7.0 kev a (cts) N H,Gal b (10 20 cm 2 ) 120 CGCG 049-057 10399 2009-04-17 VFAINT 19.06 30.2 ± 7.6 2.6 121 NGC 1068 344 2000-02-21 FAINT 47.44 100828.1 ± 326.7 2.9 123 UGC 02238 15068 2012-12-02 VFAINT 14.87 132.1 ± 13.5 8.9 127 MCG-03-34-064 7373 2006-07-31 FAINT 7.09 1029.3 ± 32.9 5.0 134 ESO 350-IG038 8175 2006-10-28 VFAINT 54.00 1794.5 ± 45.8 2.4 136 MCG -01-60-022 10570 2009-08-13 FAINT 18.90 325.4 ± 21.7 3.6 142 NGC 5135 2187 2001-09-04 FAINT 29.30 3975.9 ± 68.0 4.9 141 IC 0860 10400 2009-03-24 VFAINT 19.15 25.9 ± 7.2 1.0 147 IC 5179 10392 2009-06-21 VFAINT 11.96 555.5 ± 32.2 1.4 148 CGCG 465-012 15071 2012-12-17 VFAINT 14.87 134.0 ± 13.4 14.8 163 NGC 4418 4060 2003-03-10 FAINT 19.81 59.6 ± 15.3 1.9 169 ESO 343-IG013 15073 2013-06-13 VFAINT 14.78 139.6 ± 13.9 2.8 170 NGC 2146 3135 2002-11-16 FAINT 10.02 2144.2 ± 50.4 7.1 174 NGC 5653 10396 2009-04-11 VFAINT 16.52 387.1 ± 22.8 1.3 178 NGC 4194 7071 2006-09-09 FAINT 35.50 2410.3 ± 51.4 1.5 179 NGC 7591 10264 2009-07-05 FAINT 4.88 26.3 ± 6.1 5.6 182 NGC 0023 10401 2008-10-27 VFAINT 19.45 753.1 ± 31.9 3.4 188 NGC 7552 7848 2007-03-31 FAINT 5.08 832.8 ± 30.2 1.2 191 ESO 420-G013 10393 2009-05-13 VFAINT 12.42 759.0 ± 29.2 2.1 194 ESO 432-IG006 15074 2013-06-24 VFAINT 16.05 280.7 ± 20.0 19.3 195 NGC 1961 10531 2009-05-08 VFAINT 32.83 723.3 ± 40.0 8.1 196 NGC 7752/3 10569 2009-08-30 FAINT 11.99 96.0 ± 12.7 5.4 198 NGC 1365 6869 2006-04-20 FAINT 15.54 4644.2 ± 72.7 1.3 199 NGC 3221 10398 2009-03-19 VFAINT 19.03 323.5 ± 28.3 1.9 201 NGC 0838 15667 2013-07-21 VFAINT 58.34 1996.0 ± 49.6 2.6 Notes: (a) The source counts are corrected for background and measured in the 0.5 7.0 kev band. The counts from separate components in a single system are summed together. (b) The Galactic absorption column density is taken from the LAB HI map by Kalberla et al. (2005). 4.2.1 Hard band fitting In hard band, as emission is dominated either by an AGN component or by HMXB, explained in section 1.2, fitting is done with a powerlaw model. A correction for galactic absorption is added. A powerlaw model gives dn/de E Γ cm 2 s 1 kev 1, with the photon index Γ related to the energy index α (F E E α ) as Γ = α + 1. The photon index is determined for data with more than 30 cts in the energy range of interest. In cases where the quality of the data does not allow for a good fit, the photon index is fixed at Γ = 2.0 and only the normalization is left as a free parameter. In cases where the hard band spectra clearly corresponds to that of an absorbed AGN, intrinsic absorption is added to the powerlaw model, as a free parameter, and the photon index is fixed to Γ = 1.8. An example of such an absorbed AGN can be seen in Fig. 6. 4.2.2 Soft band fitting In the soft X-ray band, heavy blending of emission lines makes it difficult to distinguish between photoionized gas contribution form an AGN component and thermal 16

Figure 6: Chandra spectra of NGC 3256 (left), a typical starburst, and MCG-03-34-064 (right), an absorbed AGN. Both spectra are shown in the range 0.4-7 kev, including the hard and soft bands. gas emission from a starburst. However, it is generally assumed (Iwasawa et al., 2011) that the dominating component in LIRGs is the thermal emission from hot plasma, arising from Bremsstrahlung electron interaction, even if for some objects there may be an important AGN contamination. A standard thermal emission spectrum model, with a solar abundance pattern, MEKAL (Mewe et al., 1995), is used. MEKAL model with only one component does not provide a good fit, even if corrected for galactic absorption, in all cases where the quality of data is good enough to compare with other options. In those cases, two MEKAL models, one with an associated lower temperature ( 1 kev) and another with a higher temperature ( 1.5 kev), are superposed. The lower temperature component arises from a more extended plasma, easily absorbed in the inner regions of the galaxy (as its emission peaks in a softer band), resulting in an observed emission that is generated in the outer regions of the centre of the galaxy. Therefore, no intrinsic absorption from the source s gas is considered. The higher temperature component produces an emission that peaks clearly in the centre of the nucleus, and even if it penetrates the surrounding material (as emission is harder), it is heavily absorbed by it. Therefore, this component has an intrinsic absorption added as a free parameter. 4.2.3 Special cases For a few specific galaxies the hard and soft band fitting described above produced unsatisfactory results, and a single model for all the spectra (0.5 7 kev) resulted in a better adjustment. It is the case of MCG+12-02-001 (W), CGCG 049-057, NGC 4418, ESO 343-IG013 (N) and ESO 343-IG013 (S), for which a powerlaw was used to described the whole X-ray range (0.5 7 kev); and for ESO 440-IG058 (N), for which a one component MEKAL model was used for the whole range. 17

4.3 Results Results of the data analysis of Chandra spectra are presented in Table 3. For each galaxy, count rates in soft and hard band, X-ray colour, estimated X-ray fluxes and luminosities in the two bands, and the logarithmic ratio to the infrared luminosity in the 8 1000 µm (as obtained from Armus et al., 2009), are listed. Count rates are computed for each component inside a system separately, as opposed to the source counts in Table 2, as well as with flux and luminosity in each band. IR luminosity is divided into the respective components with a ratio estimated from Spitzer MIPS photometry, with the exception of the three systems marked in Table 3, as they appeared as unresolved sources. The X-ray colour is defined as HR = (H S)/(H + S), also referred to as Hardness Ratio, where H and S are background corrected counts in the 2 8 kev and in the 0.5 2 kev bands respectively. 4.3.1 AGN selection The hardness ratio gives the relative strength of the X-ray emission above and below 2 kev, in counts. Emission above 2 kev being much stronger, implying a higher HR, is often associated with an absorbed X-ray source, which indicates the presence of an obscured AGN. As described in Iwasawa et al. (2009), AGN can be selected by the hardness of the X-ray emission, classifying objects with HR > 0.3 as AGN. This threshold is chosen because ULIRGs known to host AGN cluster just above this value. This method may miss some Compton-thick AGN, as they are also absorbed in the hard band and only reflected radiation is observed. The values of HR as a function of L X are plotted in Fig. 7, the median of the distribution is -0.59. Figure 7: Hardness ratio as a function of X-ray luminosity, derived for single components in every system (see Table 3.). The threshold for selecting AGN, HR > 0.3, is shown with a dashed line. The median is -0.59. Selection of AGN can also be carried out using infrared observations. For the 18

Table 3: X-ray spectral properties for the sample. No. Galaxy SX (1) HX (2) HR (3) F SX (4) F HX (5) L SX (6) L HX (7) SX/IR (8) HX/IR (9) 45 UGC 08387 15.47 ± 1.06 4.28 ± 0.59 0.57 ± 0.04 5.18 5.57 7.50 8.07 4.44 4.41 47 CGCG 436 030 9.01 ± 0.84 3.18 ± 0.54 0.48 ± 0.04 4.07 5.50 8.75 11.83 4.33 4.20 49 NGC 0695 15.58 ± 1.06 5.59 ± 0.66 0.47 ± 0.03 6.83 11.68 15.78 27.01 4.08 3.84 50 CGCG 043 099 2.60 ± 0.44 0.92 ± 0.32 0.48 ± 0.09 1.09 1.35 3.98 4.96 4.67 4.57 51 MCG+07 23 019 8.51 ± 0.45 1.17 ± 0.24 0.76 ± 0.04 3.65 1.93 10.91 5.76 4.17 4.44 53 UGC 02369 (S) 11.16 ± 1.14 1.30 ± 0.48 0.79 ± 0.08 4.31 1.72 9.53 3.81 4.28 4.67 54 NGC 1614 38.57 ± 1.59 12.19 ± 0.91 0.52 ± 0.02 1.82 1.87 1.00 1.03 5.23 5.22 56 NGC 5331 (N) 1.79 ± 0.38 0.73 ± 0.26 0.42 ± 0.10 1.01 1.32 2.89 3.79 4.18 4.06 56 NGC 5331 (S) 4.37 ± 0.59 1.36 ± 0.36 0.53 ± 0.07 2.25 1.97 6.46 5.68 4.31 4.37 57 IRAS F06076 2139(N) 2.54 ± 0.44 1.00 ± 0.33 0.43 ± 0.09 1.37 1.23 4.45 4.01 4.59 4.63 60 IC 2810(NW) 3.88 ± 0.57 0.55 ± 0.35 0.75 ± 0.12 1.88 0.96 5.53 2.83 (a) (a) 60 IC 2810 (SE) 1.75 ± 0.39 0.12 ± 0.19 0.87 ± 0.22 0.75 0.01 2.20 0.02 (a) (a) 63 IRAS 18090+0130 (E) 4.01 ± 0.55 1.27 ± 0.36 0.52 ± 0.08 1.90 2.58 4.08 5.54 4.62 4.49 63 IRAS 18090+0130(W) 1.08 ± 0.30 0.33 ± 0.23 0.54 ± 0.18 (b) (b) (b) (b) (b) (b) 64 III Zw 035 (N) 2.34 ± 0.41 0.84 ± 0.26 0.47 ± 0.09 1.07 1.33 1.81 2.25 4.97 4.87 64 III Zw 035 (S) 0.96 ± 0.27 0.73 ± 0.24 0.14 ± 0.10 0.39 1.39 0.66 2.35 (c) (c) 65 NGC 3256 263.36 ± 3.28 28.62 ± 1.67 0.80 ± 0.01 89.99 43.60 16.29 7.89 4.01 4.33 67 IRAS F16399 0937(N) 4.02 ± 0.53 1.60 ± 0.34 0.43 ± 0.06 2.06 2.56 4.03 5.02 (a) (a) 67 IRAS F16399 0937(S) 2.06 ± 0.38 1.25 ± 0.31 0.24 ± 0.07 1.42 2.10 2.79 4.12 (a) (a) 68 IRAS F16164 0746 3.91 ± 0.55 2.80 ± 0.52 0.17 ± 0.04 1.90 6.72 3.73 13.17 4.63 4.09 71 NGC 2623 4.79 ± 0.52 3.85 ± 0.49 0.11 ± 0.02 1.76 7.95 1.49 6.72 5.01 4.36 72 IC 5298 11.76 ± 0.92 3.31 ± 0.56 0.56 ± 0.04 5.92 9.66 10.03 16.36 4.18 3.97 73 IRAS 20351+2521 9.14 ± 0.89 1.69 ± 0.52 0.69 ± 0.07 5.14 2.60 14.02 7.09 4.05 4.34 75 NGC 6090 (NE) 16.71 ± 1.07 1.84 ± 0.37 0.80 ± 0.05 6.07 2.88 13.63 6.47 (a) (a) 75 NGC6090 (SW) 4.68 ± 0.57 0.56 ± 0.21 0.79 ± 0.10 1.78 1.13 4.00 2.53 (a) (a) 79 NGC 5256 (NE) 32.39 ± 1.31 7.97 ± 0.69 0.60 ± 0.02 12.17 18.06 24.23 35.95 3.28 3.11 79 NGC 5256 (SW) 20.47 ± 1.02 1.70 ± 0.32 0.85 ± 0.04 8.32 3.31 16.56 6.58 3.75 4.15 80 IRAS F03359+1523(E) 5.92 ± 0.67 1.41 ± 0.38 0.62 ± 0.07 3.14 2.34 8.68 6.48 4.20 4.32 81 ESO 550 IG025 (N) 2.23 ± 0.41 0.55 ± 0.24 0.60 ± 0.12 1.03 1.55 2.27 3.42 4.44 4.26 81 ESO 550 IG025 (S) 1.55 ± 0.35 0.76 ± 0.26 0.34 ± 0.10 0.75 1.51 1.65 3.34 4.58 4.27 82 NGC 0034 15.58 ± 1.08 6.68 ± 0.76 0.40 ± 0.03 7.56 14.26 6.40 12.07 4.27 3.99 83 MCG+12 02 001 (E) 11.12 ± 0.89 3.77 ± 0.53 0.49 ± 0.04 5.16 5.77 3.01 3.37 4.56 4.51 83 MCG+12 02 001 (W) 3.02 ± 0.46 1.49 ± 0.34 0.34 ± 0.06 2.16 2.99 1.26 1.75 (c) (c) 85 IRAS F17138 1017 7.84 ± 0.76 6.19 ± 0.71 0.12 ± 0.02 4.67 12.87 3.94 10.86 4.48 4.04 95 ESO 440 IG058 (N) 2.12 ± 0.38 1.02 ± 0.27 0.35 ± 0.08 1.15 1.28 1.72 1.92 (c) (c) 95 ESO 440 IG058 (S) 7.16 ± 0.70 0.97 ± 0.28 0.76 ± 0.08 3.73 1.54 5.60 2.31 4.23 4.61 100 NGC 7130 78.11 ± 1.44 8.00 ± 0.53 0.81 ± 0.01 26.71 13.31 16.89 8.42 3.78 4.08 104 NGC 7771 33.16 ± 1.54 12.55 ± 1.14 0.45 ± 0.02 13.96 20.28 6.26 9.09 4.09 3.93 104 NGC 7770 (NGC7771S) 8.08 ±0.73 0.34 ± 0.29 0.92 ± 0.09 3.55 1.28 1.59 0.57 4.08 4.53 105 NGC 7592 (E) 8.42 ± 0.76 1.11 ± 0.32 0.77 ± 0.07 3.68 1.81 4.94 2.43 4.09 4.40 105 NGC 7592 (W) 11.29 ± 0.88 2.04 ± 0.41 0.69 ± 0.05 4.63 3.17 6.23 4.26 3.75 3.92 106 NGC 6286 32.15 ± 1.64 1.27 ± 0.87 0.92 ± 0.05 12.44 2.58 10.93 2.27 3.79 4.47 106 NGC 6285 (NGC 6286 NW) 5.04 ± 0.64 0.45 ± 0.33 0.83 ± 0.12 1.97 0.86 1.73 0.75 4.11 4.48 107 NGC 4922 10.93 ± 0.97 2.73 ± 0.62 0.60 ± 0.05 5.08 5.93 7.49 8.74 4.09 4.02 110 NGC 3110 21.92 ± 1.28 4.73 ± 0.77 0.65 ± 0.04 11.30 8.43 8.54 6.37 4.02 4.15 114 NGC 0232 10.91 ± 0.91 2.19 ± 0.55 0.67 ± 0.06 4.95 4.35 5.37 4.72 4.29 4.35 117 MCG+08 18 013(E) 5.27 ± 0.67 2.1 ± 0.45 0.43 ± 0.06 2.49 3.57 4.07 5.85 4.32 4.16 120 CGCG 049 057 0.83 ± 0.28 0.76 ± 0.29 0.04 ± 0.09 0.40 1.26 0.20 0.64 5.63 5.13 121 NGC 1068 1975.11 ± 6.58 150.27 ± 2.03 0.86 ± 0.00 539.24 224.33 16.31 6.79 3.77 4.15 123 UGC 02238 6.94 ± 0.74 1.94 ± 0.52 0.56 ± 0.06 3.70 4.38 3.77 4.47 4.34 4.26 127 MCG 03 34 064 95.67 ± 3.73 49.50 ± 2.75 0.32 ± 0.01 43.43 158.12 35.11 127.83 3.32 2.76 134 ESO 350 IG038 26.85 ± 0.73 6.38 ± 0.43 0.62 ± 0.02 10.91 10.06 10.34 9.53 3.85 3.89 136 MCG 01 60 022 10.52 ± 0.84 6.70 ± 0.78 0.22 ± 0.02 4.53 16.82 5.42 20.13 4.12 3.55 142 NGC 5135 123.31 ± 2.14 12.39 ± 0.90 0.82 ± 0.01 44.38 19.85 19.69 8.81 3.59 3.94 19

Table 3 continued. No. Galaxy SX (1) HX (2) HR (3) F SX (4) F HX (5) L SX (6) L HX (7) SX/IR (8) HX/IR (9) 141 IC 0860 0.68 ± 0.26 0.37 ± 0.27 0.29 ± 0.17 (b) (b) (b) (b) (b) (b) 147 IC 5179 39.74 ± 2.14 6.71 ± 1.62 0.71 ± 0.04 15.79 14.40 4.99 4.55 4.13 4.17 148 CGCG 465 012 7.13 ± 0.77 1.52 ± 0.58 0.65 ± 0.08 3.10 2.46 3.29 2.62 4.27 4.37 163 NGC 4418 2.63 ± 0.59 0.38 ± 0.32 0.75 ± 0.18 0.92 0.91 0.15 0.15 5.61 5.61 169 ESO 343 IG013 (N) 2.02 ± 0.38 2.05 ± 0.38 0.01 ± 0.05 0.96 4.46 0.85 3.93 4.72 4.05 169 ESO 343 IG013 (S) 3.15 ± 0.47 1.02 ± 0.30 0.51 ± 0.08 1.84 1.24 1.62 1.10 (c) (c) 170 NGC 2146 170.25 ± 4.38 43.74 ± 2.47 0.59 ± 0.01 68.25 76.39 2.50 2.80 4.31 4.26 174 NGC 5653 20.26 ± 1.20 3.18 ± 0.70 0.73 ± 0.04 8.29 4.44 3.59 1.93 4.16 4.43 178 NGC 4194 59.99 ± 1.33 7.91 ± 0.57 0.77 ± 0.02 23.62 12.43 5.23 2.75 3.97 4.25 179 NGC 7591 4.29 ± 1.01 1.09 ± 0.74 0.59 ± 0.17 (b) (b) (b) (b) (b) (b) 182 NGC 0023 34.84 ± 1.45 3.88 ± 0.77 0.80 ± 0.03 13.99 6.07 7.12 3.09 3.85 4.22 188 NGC 7552 148.31 ± 5.56 15.63 ± 2.09 0.81 ± 0.03 56.93 24.29 3.76 1.60 4.12 4.49 191 ESO 420 G013 56.94 ± 2.19 4.17 ± 0.83 0.86 ± 0.03 25.27 5.83 7.86 1.81 3.76 4.40 194 ESO 432 IG006 (NE) 8.09 ± 0.76 1.75 ± 0.48 0.64 ± 0.06 5.76 4.46 3.82 2.96 3.78 3.89 194 ESO 432 IG006 (SW) 6.04 ± 0.66 1.81 ± 0.46 0.54 ± 0.06 4.33 4.03 2.87 2.67 3.91 3.94 195 NGC 1961 17.55 ± 0.89 4.48 ± 0.83 0.59 ± 0.03 14.94 15.77 6.22 6.57 3.85 3.83 196 NGC 7752/3 7.05 ± 0.85 0.96 ± 0.62 0.76 ± 0.11 3.02 2.24 1.96 1.45 4.36 4.49 198 NGC 1365 177.35 ± 3.63 121.5 ± 2.95 0.19 ± 0.00 68.22 364.21 2.62 13.96 4.17 3.44 199 NGC 3221 12.47 ± 1.08 4.53 ± 1.02 0.47 ± 0.04 5.01 12.22 2.59 6.31 4.26 3.88 201 NGC 0838 30.31 ± 0.75 3.91 ± 0.39 0.77 ± 0.02 15.87 6.53 5.50 2.26 3.89 4.28 Notes: Column (1): background corrected count rate in the 0.5 2 kev band in units of 10 3 ct s 1. Column (2): background corrected count rate in the 2 7 kev band in units of 10 3 ct s 1. Column (3): X-ray colour as defined by HR = (H S)/(H + S). Column (4): observed 0.5 2 kev band flux corrected for Galactic absorption in units of 10 14 erg s 1 cm 2. Column (5): observed 2 7 kev band flux corrected for Galactic absorption in units of 10 14 erg s 1 cm 2. Column (6): 0.5 2 kev band luminosity corrected for Galactic absorption in units of 10 40 erg s 1. Column (7): 2 7 kev band luminosity corrected for Galactic absorption in units of 10 40 erg s 1. Column (8): logarithmic luminosity ratio of the 0.5 2 kev and 8 1000 µm bands. Column (9): logarithmic luminosity ratio of the 2 8 kev and 8 1000 µm bands. (a) Missing values due to inability to separate IRAS IR flux into two components (unresolved sources in MIPS 24). (b) Missing values due to inability to fit X-ray spectra, not enough source counts. (c) Missing values due to all IRAS IR flux in the system being originated in the source s companion. GOALS sample two main approaches have been followed: detection of the [Ne v] 14.3 µm line over kpc scales, which provides direct evidence for the presence of an AGN (Petric et al, 2011), and the equivalent width of the 6.2 µm PAH feature being lower than 0.1 µm (Stierwalt et al., 2013). From the data presented in the two mentioned papers and the X-ray HR selection from this work s study, which are complementary, a list of AGN within this work s sample is presented in Table 4. A total of 24 galaxies fit into at least one of the three criteria (a 32 % of our sample). 4.3.2 X-ray luminosities and correlation with IR The distribution of X-ray luminosities, compared to the Iwasawa et al. (2011) sample, is shown in Fig. 8. It has a median logarithmic value of log(l X ) 41.1 erg s 1. Iwasawa et al. sample has a median of log(l X ) 41.6 erg s 1, but a wider distribution. A weak correlation between IR and X-ray luminosity can be seen in Fig. 9, with a typical spread of over one order of magnitude. When multiple components are present in a system they are represented separately, except for the cases 20

Table 4: AGN selection according to any of the three mentioned methods in section 4.3.1. No. Galaxy [Ne v] 14.3 µm (1) 6.2 µm EQW(σ) (µm) (2) 45 UGC 08387 Y 0.62 (0.01) -0.57 ± 0.04 64 III Zw 035 (S) N (a) -0.14 ± 0.10 67 IRAS F16399-0937(S) N 0.43 (0.01) -0.24 ± 0.07 68 IRAS 16164-0746 Y 0.61 (0.01) -0.17 ± 0.04 71 NGC 2623 Y 0.27 (0.01) -0.11 ± 0.02 72 IC 5298 Y 0.12 (0.004) -0.56 ± 0.04 79 NGC 5256 (NE) Y 0.44 (0.02) (b) -0.60 ± 0.02 79 NGC 5256 (SW) Y 0.44 (0.02) (b) -0.85 ± 0.04 85 IRAS F17138-1017 N 0.68 (0.01) -0.12 ± 0.02 100 NGC 7130 Y 0.30 (0.01) -0.81 ± 0.01 105 NGC 7592 (E) Y (b) 0.48 (0.01) -0.77 ± 0.07 105 NGC 7592 (W) Y (b) 0.30 (0.01) -0.69 ± 0.05 107 NGC 4922 Y 0.16 (0.003) -0.60 ± 0.05 114 NGC 0232 Y 0.16 (0.005) -0.67 ± 0.06 120 CGCG 049-057 N 0.51 (0.04) -0.04 ± 0.09 121 NGC 1068 Y (c) -0.86 ± 0.00 127 MCG -03-34-064 Y <0.01 ( ) -0.32 ± 0.01 136 MCG -01-60-022 N (a) -0.22 ± 0.02 141 IC 0860 N 0.43 (0.01) -0.29 ± 0.17 142 NGC 5135 Y 0.49 (0.01) -0.82 ± 0.01 163 NGC 4418 N <0.02 ( ) -0.75 ± 0.18 169 ESO 343-IG013 (N) N 0.47 (0.01) 0.01 ± 0.05 191 ESO 420-G013 Y 0.30 (0.003) -0.86 ± 0.03 198 NGC 1365 N 0.13 (0.002) -0.19 ± 0.00 HR (3) Notes: (1) Detection of the [Ne v] 14.3 µm line over kpc scales in Petric et al (2011) observations. (2) Equivalent width of the 6.2 µm PAH feature in Stierwalt et al. (2013). (3) Hardness ratio as calculated with this work s data. All galaxies of this work s sample that are selected by any of the three criteria described are listed. (a) Not analysed galaxies due to archival SL staring mode observations not centred on the galaxy s nucleus. (b) Galaxies not resolved in MIR observations in any of the two mentioned papers. (c) Galaxy observation saturated the spectrograph, not analysed. where IR sources were not resolved (shown in Table 3), being then plotted with the system s overall luminosity. The data shows a significant spread, caused by the scatter around the linear correlation, rather than any non-linear one. The adjusted correlation has slope 1 in the logarithmic plane, for better comparison with results from previous works. Adjustment is log(l SX ) = log(l ir ) 4.22±0.44 for the soft X-rays and log(l HX ) = log(l ir ) 4.24±0.42 for the hard X-rays, where the uncertainties are the dispersion of the distribution. Correlation coefficients for the two distributions are, respectively, 0.24 and 0.37. Values obtained by Iwasawa et al. in their higher IR luminosity GOALS subsample are log(l SX ) = log(l ir ) 4.53±0.34 and log(l HX ) = log(l ir ) 4.40±0.63. Removing sources indicated as AGN in Table 4 from our sample results in obtaining log(l SX ) = log(l ir ) 4.21 ± 0.28 for the soft band and log(l HX ) = log(l ir ) 4.29 ± 0.30 for the hard band. Correlation coefficients then become 0.26 and 0.46, respectively. Comparison between the obtained results at high star formation rates for this work s sample and Iwasawa et al. (2011) sample are shown 21